The process by which heavier chemical elements are synthesized from hydrogen nuclei in the interiors of stars.
nucleosynthetic nu'cle·o·syn·thet'ic (-sĭn-thĕt'ĭk) adj.
nucleosynthesis
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Theories of the origin of the elements involve synthesis with charged and neutral elementary particles (neutrons, protons, neutrinos, photons) and other nuclear building blocks of matter, such as alpha particles. The theory of nucleosynthesis comprises a dozen distinct processes, including big bang nucleosynthesis, cosmic-ray spallation in the interstellar medium, and static or explosive burning in various stellar environments (hydrogen-, helium-, carbon-, oxygen-, and silicon-burning, and the, s-, r-, p-, γ-, and ν-processes). Acceptable theories must lead to an understanding of the cosmic abundances observed in the solar system, stars, and the interstellar medium. Hydrogen and helium constitute about 98% of the total element content by number of atoms, and there is a rapid decrease with increasing nuclear mass number A. See also Elements, cosmic abundance of.
Observations of the expanding universe and of the 3-K background radiation indicate that the universe originated in a primordial event known as the big bang about 15 × 109 years ago. Absence of stable mass 5 and mass 8 nuclei preclude the possibility of synthesizing the major portion of the nuclei of masses greater than 4 in those first few minutes of the universe, when density and temperature were sufficiently high to support the necessary nuclear reactions to synthesize elements. See also Big bang theory; Cosmic background radiation;
The principal source of energy in stars is certainly nuclear reactions which release energy by fusion of lighter nuclei to form more massive nuclei. There is considerable evidence that nucleosynthesis has been going on in stars for billions of years. Observations show that the abundance ratio of iron and heavier elements to hydrogen decreases with increasing stellar age. The oldest known stars in the disk of the Milky Way Galaxy exhibit a ratio 10,000 times smaller than in the Sun. This low ratio is understood on the basis of element synthesis occurring in previous-generation stars that evolved to the point of exploding as supernovae, thus enriching the interstellar medium with the nuclei that were synthesized during their lifetimes and by the explosive synthesis that accompanies the ejection of the stellar envelope. Later-generation stars were then formed from enriched gas and dust. See also Milky Way Galaxy;
Hydrogen burning, the first process of nucleosynthesis, converts hydrogen to helium. In stars of 1.2 or less solar masses, this process occurs via the proton-proton chain. This chain was also responsible for much of the element synthesis during the big bang, producing the bulk of deuterium and helium, and some of the 7Li, observed today. In more massive stars where the central temperatures exceed 2 × 107 K, hydrogen burning is accomplished through proton captures by carbon, nitrogen, and oxygen nuclei, in the carbon-nitrogen-oxygen (CNO) cycles, to form 4He. The product (ash) of hydrogen burning is helium, but much of the helium produced is consumed in later stages of stellar evolution or is locked up forever in stars of lower mass that never reach the temperatures required to ignite helium burning. The observed abundances of some carbon, nitrogen, oxygen, and fluorine nuclei are attributed to hydrogen burning in the CNO cycles. See also Carbon-nitrogen-oxygen cycles;
When the hydrogen fuel is exhausted in the central region of the star, the core contracts and its temperature and density increase. Helium, the ash of hydrogen burning, cannot be burned immediately due to the larger nuclear charge of helium (Z = 2) producing a much higher Coulomb barrier against fusion. When the temperature eventually exceeds about 108 K, helium becomes the fuel for further energy generation and nucleosynthesis. The basic reaction in this thermonuclear phase is the triple-alpha process in which three 4He nuclei (three alpha particles) fuse to form 12C, a carbon nucleus of mass 12 (atomic mass units). Capture of an alpha particle by 12C then forms 16O, symbolically written as 12C + 4He → 16O + γ, or simply 12C(α,γ)16O, where γ represents energy released in the form of electromagnetic radiation. Other reactions that are included in helium burning are 16O(α,γ)20Ne, 20Ne(α,γ)24Mg, 14N(α,γ)18F, and 18O(α,γ)22Ne. Fluorine-18, produced when 14N captures an alpha particle, is unstable and decays by emitting a positron (e+) and a neutrino (ν) to form 18O [in short, 14N(α,γ)18F(e+,ν)18O]. Because there is likely to be 13C in the stellar core if hydrogen burning proceeded by the carbon-nitrogen-oxygen cycles, the neutron-producing reaction 13C(α,n)16O should also be included with the helium-burning reactions. Helium burning is probably responsible for much of the 12C observed in the cosmic abundances, although in more massive stars the later burning stages will consume the 12C produced earlier by helium burning. See also
Upon exhaustion of the helium supply, if the star has an initial mass of at least 8 solar masses, gravitational contraction of the stellar core can lead to a temperature exceeding 5 × 108 K, where it becomes possible for two 12C nuclei to overcome their high mutual Coulomb-repulsion barrier and fuse to form 20Ne, 23Na, and 24Mg through reactions such as 12C(12C,α)20Ne, 12C(12C,p)23Na, and 12C(12C,γ)24Mg. Carbon burning can produce a number of nuclei with masses less than or equal to 28 through further proton and alpha-particle captures.
Carbon burning is followed by a short-duration stage, sometimes referred to as neon burning, in which 20Ne disintegrates by the reaction 20Ne(γ,α)16O. The eventual result is that most of the carbon from helium burning becomes oxygen, which supplements the original oxygen formed in helium burning. This stage is followed by the fusion of oxygen nuclei at much higher temperatures. (Temperatures greater than 109 K are required for 16O nuclei to overcome their mutual Coulomb barrier.) Some relevant reactions for oxygen burning are 16O(16O,α)28Si, 16O(16O,p)31P, and 16O(16O,γ)32S. Nuclei of masses up to A = 40 may be produced in this phase through proton, neutron, and alpha-particle captures.
Silicon burning commences when the temperature exceeds about 3 × 109 K. In this phase, photodisintegration of 28Si and other intermediate-mass nuclei around A = 28 produces copious supplies of protons, neutrons, and alpha particles. These particles capture on the seed nuclei left from previous burning stages and thus produce new isotopes up to mass 60, resulting in the buildup of the abundance peak near A = 56.
Because neutrons are neutral particles, their capture is not affected by the Coulomb barrier that inhibits charged-particle reactions. If the number of neutrons per seed nucleus is small, so that time intervals between neutron captures are long compared to the beta-decay lifetimes of unstable nuclei that are formed, the s-process (slow process) takes place. The seed nuclei are predominantly in the iron peak, but the abundances of low-mass nuclei are also affected by neutron-capture processing. In the s-process, if neutron capture produces a nucleus that is unstable (due to an excess number of neutrons), the newly formed nucleus undergoes beta decay to a stable isobar by emitting an electron and an antineutrino. The resulting nucleus eventually captures another neutron, and the capture-decay step repeats. The presence of free neutrons leads to a chain of capture-decay events that drives the abundances along a unique s-process path that zigzags along a single line in the nuclear Z-N diagram near the low of beta-stable nuclei (valley of beta stability). Given enough neutrons, the s-process synthesizes nuclei of masses up to 209, when alpha decay becomes a deterrent to further buildup by neutron capture. See also Radioactivity.
The r-process occurs when a large neutron flux allows rapid neutron capture, so that seed nuclei capture many neutrons before undergoing beta decay. The rapid neutron capture takes the nuclei far away from the valley of beta stability, into the regime of extremely neutron-rich nuclei. The time scale for the r-process is very short, 1–100 s, and the abundance of free neutrons is very large. These conditions are found deep inside the interiors of exploding massive stars, supernovae. The r-process can synthesize nuclei all the way into the transuranic elements.
The p-process produces heavier elements on the proton-rich side of the beta valley. The p-nuclei (about 35 are known) are blocked from formation by stable nuclei produced by either the r- or s-process. The major task for p-process theory is thus to find ways, other than beta decay, to process the more abundant r- and s-nuclei into the less abundant p-nuclei. Two possible mechanisms are: radiative proton capture, and gamma-induced neutron, proton, or alpha-particle removal reactions. In both cases the temperature should be in excess of 2–3 × 109 K.
The neutrino (ν) flux emitted from a cooling proto neutron star alters the yields of explosive nucleosynthesis from type II supernovae. Inelastic scattering of neutrinos (all flavors) off abundant nuclei excites states that can decay via single or multiple nucleon emission. The ν-process is probably responsible for significant contributions to the synthesis in nature of about a dozen isotopes. While the neutrino interaction cross section with matter is extremely small (about 10−44 cm2), the high neutrino energies and the large number flux close to the collapsing iron core of a massive star lead to significant synthesis of nuclei, either directly or indirectly. See also Neutrino.
The bulk of the light elements lithium, beryllium, and boron found in the cosmic abundance curve cannot have survived processing in stellar interiors because they are readily destroyed by proton capture. Although some 7Li originated in the big bang, primordial nucleosynthesis cannot be responsible for the bulk of the 7Li in existence. Spallation of more abundant nuclei such as carbon, nitrogen, and oxygen by high-energy protons and alpha particles can account for the low-abundance nuclides 6Li, 9Be, 10B, 11B, and for some 7Li. The canonical process is spallation of carbon, nitrogen, and oxygen nuclei in the interstellar medium by fast light particles, such as alpha particles and protons, which are abundant in the gas between the stars. These high-energy particles are referred to as cosmic rays, leading to the term “cosmic-ray spallation (CRS) process.” See also Cosmic rays.
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nucleosynthesis or nucleogenesis, in astronomy, production of all the chemical elements from the simplest element, hydrogen, by thermonuclear reactions within stars, supernovas, and in the big bang at the beginning of the universe (see nucleus; nuclear energy). A star obtains its energy by fusing together light nuclei to form heavier nuclei; in this process, mass (m) is converted into energy (E) in accordance with Einstein's formula, E=mc2, in which c is the speed of light. The reactions are initiated by the high temperatures (about 14 million degrees Celsius) at the center of the star. In the course of producing nuclear energy, the star synthesizes all the elements of the periodic table from its initial composition of mostly hydrogen and a small amount of helium.
Transformation of Hydrogen to Helium
The first step is the fusion of four hydrogen nuclei to make one helium nucleus. This “hydrogen-burning” phase supplies energy to stars on the main sequence of the Hertzsprung-Russell diagram. There are two chains of reactions by which the conversion of hydrogen to helium is effected: the proton-proton cycle and the carbon-nitrogen-oxygen cycle (sometimes referred to simply as the carbon cycle). They were both first studied and proposed as sources of stellar energy by H. Bethe and independently by C. von Weiszäcker. The proton-proton cycle operates in less massive and luminous stars like the sun, while the carbon-nitrogen-oxygen cycle (which speeds up dramatically at higher temperatures) dominates in more massive and luminous stars.
The Proton-Proton Cycle
In the proton-proton cycle, two hydrogen nuclei (protons) are fused and one of these protons is converted to a neutron by beta decay (see radioactivity) to make a deuterium nucleus (one proton and one neutron). Then a third proton is added to deuterium to form the light isotope of helium, helium-3. When two helium-3 nuclei collide, they form a nucleus of ordinary helium, helium-4 (two protons and two neutrons), and release two protons. In each of these steps considerable energy is also released.
The Carbon-Nitrogen-Oxygen Cycle
The carbon-nitrogen-oxygen cycle requires minute traces of carbon as a catalyst. Four protons are added, one by one, to a carbon nucleus to form a succession of excited (unstable) nuclei of carbon, nitrogen, and oxygen. The intermediate nuclei shed their excess electric charge via beta decay and the final oxygen nucleus spontaneously splits into the original carbon nucleus and a helium-4 nucleus, releasing energy. The net effect is again the combination of four hydrogen nuclei to form one helium-4 nucleus; the carbon is free to begin the cycle over again.
Creation of the Heavier Elements
After the bulk of a star's hydrogen has been converted to helium by either the proton-proton or carbon-nitrogen-oxygen process, the stellar core contracts (while the outer layers expand) until sufficiently high temperatures are reached to initiate “helium-burning” by the triple-alpha process; in this process, three helium nuclei (alpha particles) are fused to make a carbon nucleus. By successive additions of helium nuclei, the heavier elements through iron-56 are built up. The elements whose atomic weights are not multiples of four are created by side reactions that involve neutrons. Because iron-56 is the most stable of the elements, it is very difficult to add an extra helium nucleus to it. However, iron-56 will readily capture a neutron to form the less stable isotope, iron-57. From iron-57, the elements through bismuth-209 can be synthesized. The elements more massive than bismuth-209 are radioactive; that is, they spontaneously break apart. However, during a supernova, an extremely intense flux of neutrons is generated and nuclear reactions proceed so rapidly that the radioactive elements do not have enough time to decay, resulting in the rapid creation of the radioactive elements up to and beyond uranium.
Bibliography
See D. L. Clayton, Principles of Stellar Evolution and Nucleosynthesis (1968, repr. 1983).
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Nucleosynthesis is the process of creating new atomic nuclei from preexisting nucleons (protons and neutrons). It is thought that the primordial nucleons themselves were formed from the quark-gluon plasma from the Big Bang as it cooled below two trillion degrees. A few minutes afterward, starting with only protons and neutrons, nuclei up to lithium and beryllium (both with mass number 7) were formed but only in relatively small amounts. This first process of primordial nucleosynthesis may also be called nucleogenesis. The subsequent nucleosynthesis of the elements (including all carbon, all oxygen, etc.) occurs primarily in stars, either by nuclear fusion (including neutron capture) or nuclear fission.
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History
The first ideas were that the chemical elements were created at the beginnings of the universe, but no successful picture could be found. Hydrogen and helium were clearly far more abundant than any of the other elements (all the rest of which constituted less than 2% of the mass of the solar system). At the same time it was clear that carbon was the next most common element, and also that there was a general trend toward abundance of light elements, especially those composed of even numbers of helium-4 nuclei.
Arthur Stanley Eddington first suggested in 1920 that stars obtain their energy by fusing hydrogen to helium, but this idea was not generally accepted because it lacked nuclear mechanisms. In the years immediately before World War II Hans Bethe first provided those nuclear mechanisms by which hydrogen is fused into helium. However, neither of these early works on stellar power addressed the origin of the elements heavier than helium. Fred Hoyle's original work on nucleosynthesis of heavier elements in stars occurred just after World War II.[1] This work attributed production of heavier elements from hydrogen in stars during the nuclear evolution of their compositions. Hoyle's work explained how the abundances of the elements increased with time as the galaxy aged. Subsequently, Hoyle's picture was expanded during the 1960s by creative contributions by William A. Fowler, Alastair G. W. Cameron, and Donald D. Clayton, and then by many others. The creative 1957 review paper by E. M. Burbidge, G. R. Burbidge, Fowler and Hoyle (see Ref. list) is a well-known summary of the state of the field in 1957. That paper defined new processes for changing one heavy nucleus into others within individual stars, processes that could be documented by astronomers.
Processes
There are a number of astrophysical processes which are believed to be responsible for nucleosynthesis in the universe. The majority of these occur within the hot matter inside stars. The successive nuclear fusion processes which occur inside stars are known as hydrogen burning (via the proton-proton chain or the CNO cycle), helium burning, carbon burning, neon burning, oxygen burning and silicon burning. These processes are able to create elements up to iron and nickel, the region of the isotopes having the highest binding energy per nucleon. Heavier elements can be assembled within stars by a neutron capture process known as the s process or in explosive environments, such as supernovae, by a number of processes. Some of the more important of these include the r process, which involves rapid neutron captures, the rp process, which involves rapid proton captures, and the p process (sometimes known as the gamma process), which involves photodisintegration of existing nuclei.
The four major types of nucleosynthesis
Big Bang nucleosynthesis
Big Bang nucleosynthesis occurred within the first three minutes of the beginning of the universe and to be responsible for much of the abundance ratios of 1H (protium), 2H (deuterium), 3He (helium-3), and 4He (helium-4), in the universe [2]. Although 4He continues to be produced by other mechanisms (such as stellar fusion and alpha decay) and trace amounts of 1H continue to be produced by spallation and certain types of radioactive decay (proton emission and neutron emission), most of the mass of these isotopes in the universe, and all but the insignificant traces of the 3He and deuterium in the universe produced by rare processes such as cluster decay, are thought to have been produced in the Big Bang. The nuclei of these elements, along with some 7Li, and 7Be are believed to have been formed when the universe was between 100 and 300 seconds old, after the primordial quark-gluon plasma froze out to form protons and neutrons. Because of the very short period in which Big Bang nucleosynthesis occurred before being stopped by expansion and cooling, no elements heavier than lithium could be formed. (Elements formed during this time were in the plasma state, and did not cool to the state of neutral atoms until much later).
Stellar nucleosynthesis
Stellar nucleosynthesis occurs in stars during the process of stellar evolution. It is responsible for the generation of elements from carbon to iron by nuclear fusion processes. Stars are the nuclear furnaces in which H and He are fused into heavier nuclei, a process which occurs by proton-proton chain in stars cooler than the Sun, and by the CNO cycle in stars more massive than the Sun.
Of particular importance is carbon, because its formation from He is a bottleneck in the entire process. Carbon is produced by the triple-alpha process in all stars. Carbon is also the main element used in the production of free neutrons within the stars, giving rise to the s process which involves the slow absorption of neutrons to produce elements heavier than iron and nickel (57Fe and 62Ni). Carbon and other elements formed by this process are also fundamental to life.
The products of stellar nucleosynthesis are generally distributed into the universe through mass loss episodes and stellar winds in stars which are of low mass, as in the planetary nebulae phase of evolution, as well as through explosive events resulting in supernovae in the case of massive stars.
The first direct proof that nucleosynthesis occurs in stars was the detection of technetium in the atmosphere of a red giant in the early 1950s[3], prototypical for the class of Tc-rich stars. Because technetium is radioactive, with halflife much less than the age of the star, its abundance must reflect its creation within that star during its lifetime. Less dramatic, but equally convincing evidence is of large overabundances of specific stable elements in a stellar atmosphere. An historically important case was observation of barium abundances some 20-50 times greater than in unevolved stars, which is evidence of the operation of the s process within that star. Many modern proofs appear in the isotopic composition of Stardust, solid grains that condensed from the gases of individual stars and which have been extracted from meteorites. Stardust is one component of cosmic dust. The measured isotopic compositions demonstrate many aspects of nucleosynthesis within the stars from which the stardust grains condensed [4]
Explosive nucleosynthesis
This includes supernova nucleosynthesis, and produces the elements heavier than iron by an intense burst of nuclear reactions that typically last mere seconds during the explosion of the supernova core. In explosive environments of supernovae, the elements between silicon and nickel are synthesized by fast fusion. Also in supernovae further nucleosynthesis processes can occur, such as the r process, in which the most neutron-rich isotopes of elements heavier than nickel are produced by rapid absorption of free neutrons released during the explosions. It is responsible for our natural cohort of radioactive elements, such as uranium and thorium, as well as the most neutron-rich isotopes of each heavy element.
The rp process involves the rapid absorption of free protons as well as neutrons, but its role is less certain. Explosive nucleosynthesis occurs too rapidly for radioactive decay to increase the number of neutrons, so that many abundant isotopes having equal even numbers of protons and neutrons are synthesized (nuclides which consist essentially of whole numbers of helium nuclei). These are stable up to 40Ca, but heavier such nuclides include the radioactive 44Ti , 48Cr , 52Fe , and 56Ni , all of which decay after the explosion to create abundant stable isotopes at each atomic weight. Many such decays are accompanied by emission of gamma-ray lines capable of identifying the isotope that has just been created in the explosion. The most convincing proof of explosive nucleosynthesis in supernovae occurred in 1987 when gamma-ray lines were detected emerging from supernova 1987A. Gamma ray lines identifying 56Co and 57Co , whose radioactive halflives limit their age to about a year, proved that 56Fe and 57Fe were created by radioactive parents. This nuclear astronomy was predicted in 1969 [5] as a way to confirm explosive nucleosynthesis of the elements, and that prediction played an important role in the planning for NASA's successful Compton Gamma-Ray Observatory. Other proofs of explosive nucleosynthesis are found within the stardust grains that condensed within the interiors of supernovae as they expanded and cooled. Stardust grains are one component of cosmic dust. In particular, radioactive 44Ti was measured to be very abundant within supernova stardust grains at the time they condensed during the supernova expansion [6], confirming a 1975 prediction for identifying supernova stardust. Other unusual isotopic ratios within these grains reveal many specific aspects of explosive nucleosynthesis.
Cosmic ray spallation
Cosmic ray spallation produces some of the lightest elements present in the universe (though not significant deuterium). Most notably spallation is believed to be responsible for the generation of almost all of 3He and the elements lithium, beryllium and boron (some lithium-7 and beryllium-7 are thought to have been produced in the Big Bang). The spallation process results from the impact of cosmic rays (mostly fast protons) against the interstellar medium. These impacts fragment carbon, nitrogen and oxygen nuclei present in the cosmic rays, and also these elements being struck by protons in cosmic rays. The process results in these light elements (Be, B, and Li) being present in cosmic rays at much higher proportion than they are represented in solar atmospheres, whereas H and He nuclei are represented in cosmic rays with approximately primordial abundance with regard to each other.
Beryllium and boron are not significantly produced in stellar fusion processes, because the instability of any 8Be formed from two 4He nuclei prevents simple 2-particle reaction building-up of these elements.
Empirical evidence
Theories of nucleosynthesis are tested by calculating isotope abundances and comparing with observed results. Isotope abundances are typically calculated by calculating the transition rates between isotopes in a network. Often these calculations can be simplified as a few key reactions control the rate of other reactions.
See also
References
- ^ Autobiography William A. Fowler
- ^[dead link] Big Bang Java Calculator v1.1, Craig Hogan, Luis Mendoza
- ^ S. Paul W. Merrill (1952). "Spectroscopic Observations of Stars of Class S". The Astrophysical Journal 116: 21. doi:. http://adsabs.harvard.edu/abs/1952ApJ...116...21M.
- ^ D. D. Clayton and L. R. Nittler (2004). "Astrophysics with Presolar Stardust". Annual Review of Astronomy and Astrophysics 42: 39–78. doi:.
- ^ D. D. Clayton, S.A. Colgate, G.J. Fishman (1969). "Gamma ray lines from young supernova remnants". The Astrophysical Journal 155: 75–82. doi:.
- ^ D. D. Clayton, L. R.Nittler (2004). "Astrophysics with Presolar stardust". Annual Reviews of Astronomy and Astrophysics 42: 39–78. doi:.
Further reading
- E. M. Burbidge, G. R. Burbidge, W. A. Fowler, F. Hoyle, Synthesis of the Elements in Stars, Rev. Mod. Phys. 29 (1957) 547 (article at the Physical Review Online Archive (subscription required)).
- F. Hoyle, Monthly Notices Roy. Astron. Soc. 106, 366 (1946)
- F. Hoyle, Astrophys. J. Suppl. 1, 121 (1954)
- D. D. Clayton, "Principles of Stellar Evolution and Nucleosynthesis", McGraw-Hill, 1968; University of Chicago Press, 1983, ISBN 0-226-10952-6
- C. E. Rolfs, W. S. Rodney, Cauldrons in the Cosmos, Univ. of Chicago Press, 1988, ISBN 0-226-72457-3.
- D. D. Clayton, "Handbook of Isotopes in the Cosmos", Cambridge University Press, 2003, ISBN 0 521 823811.
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- What is stellar nucleosynthesis?
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- population III (astronomy)
- rapid proton capture process (astrophysics)
- Elements, cosmic abundance of (geochemistry)
- Hans Albrecht Bethe (American physicist)
- Spallation reaction (nuclear physics)
- William Alfred Fowler (American physicist)
- Geochronometry
- Cosmochemistry
- stellar populations (in astronomy)
- Carbon star (stars, the galaxy)
- Charged particle beams (physical chemistry)
- stellar structure (in astrophysics)
- Wolf-Rayet star
- oxygen (element – in chemistry)
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